EC 11481-2303 - A Peculiar Subdwarf OB Star Revisited

0.59 MB·

by  T. Rauch

AstrophysicsandSpaceScience
DOI10.1007/s•••••-•••-••••-•
EC 11481−2303 - A Peculiar Subdwarf OB Star Revisited
Thomas Rauch • Klaus Werner • Jeffrey W. Kruk
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(cid:13)c Springer-Verlag••••
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2 Abstract EC11481−2303 is a peculiar, hot, high- to 40000K and more,and they differ spectroscopically

gravity pre-white dwarf. Previous optical spectroscopy from the sdBs by the appearance of lines of ionized
]
R revealed that it is a sdOB star with Teff=41790K, helium, e.g. the most prominent Heiiλ4686˚A line.
S logg=5.84, and He/H=0.014 by number. We present The spectral analysis of sdOB stars is, thus, more reli-
h. an on-going spectral analysis by means of non-LTE ablebecauseTeff maybepreciselydeterminedfromthe
p model-atmospheretechniquesbasedonhigh-resolution, Hei/Heii ionization equilibrium.
- high-S/N optical (VLT-UVES) and ultraviolet (FUSE, Metal-abundance determinations in Extended Hor-
o
r IUE) observations. izontal Branch (EHB) stars are of particular interest,
st We are able to reproduce the optical and UV ob- not only because they might provide insight into their
a servations simultaneously with a chemically homoge- evolutionary history, but also because they may shed
[
neous NLTE model atmosphere with a significantly light on the question of pulsational instability in these

1 higher effective temperature and lower He abundance stars. SeveralsdBstarsdefineanewinstabilitystripin
v (Teff=55000K, logg=5.8, and He/H= 0.0025 by the Hertzsprung-Russell Diagram (HRD) (see, e.g. the
5
number). While C, N, and O appear less than 0.15
1 review by O’Donoghue etal. 1999), that was predicted
times solar, the iron-group abundance is strongly en-
1 from pulsational models by Charpinet etal. (1996). A
5 hanced by at least a factor of ten. high iron abundance in subphotospheric layers is re-
.
1 Keywords Stars: individual: EC11481−2303- Stars: quired for pulsation driving.
0 In an analysis of a large number of EHB stars,
0 abundances - Stars: atmospheres - Stars: chemically
Edelmann(2003)detectedthattwosdOBsandonesdB
1 peculiar - subdwarfs
: are very peculiar. From optical spectra he found ex-
v
treme overabundances for many iron-group elements,
i
X 1 Introduction up to 30000times solar,and it is thoughtthat the ori-
r gin is radiative levitation of these elements.
a Subdwarfs of spectral type B and OB (sdB and sdOB In this paper we examine the case of another pecu-
stars, respectively) represent an extension of Horizon- liar sdOB star, EC11481−2303(McCook & Sion 1999,
tal Branch B (HBB) stars towards higher effective WD1148−230) which was first analyzed by Stys etal.
temperatures. The sdB stars are found in the range
(2000) (Sect.2).
Teff≈25000–30000K. The sdOB stars are hotter, up
ThomasRauch 2 Discovery and first spectral analysis
KlausWerner
InstituteforAstronomyandAstrophysics,KeplerCenterforAs- EC11481−2303was discoveredin the Edinburgh-Cape
tro and Particle Physics, Eberhard Karls University, Sand 1, Blue Object Survey (V = 11.76,B−V = −0.27,U −
72076Tu¨bingen,Germany B = −1.16, Kilkenny etal. 1997). A faint companion
JeffreyW. Kruk was detected at a distance of 6′.′6, too far away to have
Department of Physics and Astronomy, Johns Hopkins Univer- an influence on the evolution of EC11481−2303.
sity,Baltimore,MD21218,U.S.A.
AfirstspectralanalysiswasperformedbyStys etal.
(2000). Based on optical (3350 − 5450˚A, resolution
2
3.5˚A, Aug 15, 1995, 1.9m telescope at SAAO) and ul- o
HeI 4472A
traviolet (1150− 1950˚A, resolution 7˚A (SWP48111)
and 0.1˚A (SWP48112), Jul 14, 1993, IUE1) observa- HeI 5876Ao
tions, they used LTE (Local Thermodynamic Equilib-
o
HeII 5412A
rium) model atmospheres (provided by Detlev Koester
4
andPierreBergeron)thatconsideredopacitiesofHand HeII 4686Ao
He only. They tried many model assumptions, homo-
H11 / HeII
geneously mixed, chemically stratified, “spot” models,
DA + DB binary models, and weak-wind models to
H10 / HeII
achieve the best fit to the H Balmer lines Hβ to H8 x
u3
with a homogeneous, single-star model. They deter- e fl H9 / HeII
mined Teff=41790K, logg=5.84, and He/H=0.014 ativ H8 / HeII
by number. With these parameters, no satisfactory el
fit to the IUE data was possible. Quite unexpect- r Hε / HeII
edly, the observed spectrum is not as steep as it is
expected for such a hot subdwarf, but instead it is 2 Hδ / HeII
ratherflat. Stys etal.(2000)speculatedonseveralpos- Hγ / HeII
sible explanations (composite binary spectrum, spot
model), and concluded that EC11481−2303 resembles Hβ / HeII
theDABwhitedwarfGD323. Wesuggestthatextreme Hα / HeII
lineblanketingbystronglyoverabundantiron-groupel- 1
[H] = 0.128
ementsmaybethetruereasonforthepeculiarUV-flux
[He] = -1.501
shape of EC11481−2303.
-20 0 20
λ - λ / Ao
0
3 Preliminary analysis
Fig. 1 Comparison of theoretical H and He line profiles
computed from H+He model atmospheres with observa-
Forouranalysis,weusedmuchbetterobservations(re- tions. The profilescalculated with parameters of Stysetal.
solution0.1˚A) inthe opticalwavelengthrange(3300˚A (2000,Tab.1)areshowninthin,dashed,bluelines,Profiles
< λ < 7000˚A) which are provided in the frame- withourmodelparametersareoverplottedinthick,full,red
work of the ESO2 SupernovaeIa Progenitor surveY lines. [X] denotes log (mass fraction / solar mass fraction)
(SPY, Napiwotzki et al. 2001) with the VLT3. We of species X.
used TMAP4, the Tu¨bingen NLTE Model-Atmosphere
Package (Rauch & Deetjen 2003; Werner et al. 2003) In a first step, we use the H and He lines to check
for the calculationofplane-parallelmodel atmospheres for the result of Stys etal. (2000). We get significantly
in hydrostatic and radiative equilibrium. We have to different results (Tab.1). This is because, based on
note here, thatthe considerationof Hiline-broadening the new observations, we can precisely evaluate the
in the calculation of spectral energy distributions Hei/Heiiionizationequilibriumandthenmeasurethe
(SEDs) has changed in TMAP. First, Repolust etal. H / He abundance ratio.
(2005) found an error in the Hi line-broadening ta-
bles (for high members of the spectral series only) by
Lemke (1997) that had been used before by TMAP. Table1 BasicparametersofEC11481−2303,derivedfrom
H+Hemodel atmospheres.
TheseweresubstitutedbyaHoltsmarkapproximation.
Second, Tremblay & Bergeron (2009) presented new, Stysetal. (2000) ourwork
parameter-free Stark line-broadening tables for Hi in- Teff / K 41790 55000
logg / (cm/sec2) 5.84 5.8
cluding non-ideal effects. These replace Lemke’s data
H / He(mass) 18 100
for the lowest ten members of the Lyman and Balmer
He / H (number) 0.014 0.0025
series.
In Fig.1, we show a comparison of spectral lines,
1International UltravioletExplorer calculated with the parameters from Tab.1 with ob-
2EuropeanSouthernObservatory servations. The theoretical line profiles of the lower
3VeryLargeTelescope members of the H Balmer series strongly deviate from
4http://astro.uni-tuebingen.de/∼rauch/TMAP.html observations. The reason is the so-called Balmer line
EC11481−2303 -APeculiarSubdwarfOBStarRevisited 3
o
HeI 4472A
200 [[HH]e]==-01..142075 -01..142075 oII 4686A oI 5876AoI 4471A HeI 5876Ao
e ee
[C] = -0.873 H HH o
HeII 5412A
[N] = -0.873
[O] = -0.873 4
K150 HeII 4686Ao
k
T / αβγδε89 H11 / HeII
HHHHHHH
100
H10 / HeII
formation depth of line core:
3
H+He H9 / HeII
50 x
H+He+C+N+O Spatyras mete tael.rs e flu H8 / HeII
v
-7 -6 -5 -4 -3 -2 -1 0 1 2 ati Hε / HeII
el
log m / g/cm2 r2 Hδ / HeII
Fig. 2 Photospheric temperaturestratification ofH+ He
Hγ / HeII
models(blue)comparedtoaH+He+C+N+Omodels
(red). The formation depths of H and He line cores are Hβ / HeII
marked.
Hα / HeII
1
problem (Werner 1996; Bergeron etal. 1993). The ne- [H] = 0.128
[He] = -1.501
glectionofmetalopacitiesinthemodel-atmospherecal-
[C] = -0.875
culation yields an incorrect temperature stratification [N] = -0.875
intheline-formingregion. Fig.2showsacomparisonof [O] = -0.875
[IG] = 0.000
thetemperaturestructureofH+HetoH+He+C+N+O
modelatmospheres(theopticalspectrumprovidesonly -20 0 20
λ - λ / Ao
upper limits for the C, N, and O abundances, we 0
adopted these values). The effect is stronger for the
Fig. 3 Same as Fig.1. In our models we consider in ad-
lower Balmer-series members. In Fig.2, we show also
dition opacities of C, N, O, and of the iron-group (IG) ele-
the temperature structure of a model that was calcu-
ments.
lated with the parameters of Stys etal. (2000, Tab.1).
DuetothemuchlowerTeff,itstemperatureisalsomuch an analysis of the UV spectrum of EC11481−2303 in
lower in the line-forming regions. This lets the fits of
the following section.
theBalmerlinescalculatedfromthismodelappearbet-
ter than the fits of our model although the Stys etal.
(2000) parameters are unrealistic.
4 UV: observations and spectral analysis
In order to demonstrate the impact of metal opaci-
ties, we include C, N, O, and the iron group (IG, here
Ca – Ni) in our TMAP models. For C, N, and O, we In addition to the IUE observations that were already
used by Stys etal. (2000), we use FUSE5 observations
usedtheupperabundancelimits (seeabove),forIGwe
(May 21, 2001, obs id B0540901). They were per-
assumeasolarmassfraction(Asplund et al.2005)). In
formed in four groups (exposure times 432−567sec)
Fig.3,wecomparetheoreticalandobservedlineprofiles
of Hi, Hei, and Heii in the optical wavelength range. that were separatedby about one FUSE orbitalperiod
(≈6000sec). The total exposure time is 8300sec. The
The agreement of the line wings of the lower members
S/Nratiois30−50perpixelataresolutionof≈0.07˚A.
of the H Balmer series is much better compared to the
Wecannotdetectanysystematicfluxvariationinthese
H+He models (Fig.2).
fourgroups. Thus,thereisnodirecthintforacompan-
We can summarize, that the optical spectrum of
EC11481−2303canbe reproducedby achemically not ion star.
stratifiedNLTE modelatmospherewithTeff=55000± The UV observationsareshowninFig.4. Their flux
5000K,logg=5.8±0.3,andH/He=100±0.3dex(by levels match well. The comparison to TMAP SEDs
showsclearly,thatreddeningcannotbetheonlyreason
mass) when metal additional opacities are considered.
Since the opticalspectrumdoes not providefurther in-
formation about metal opacities, we will continue with 5FarUltravioletSpectroscopic Explorer
4
0
FUSE IUE IGVI
(Gaussian 7Ao FWHM) IGV
-5
o-1A n IGVII
-1-2secgcm1 EB-V == 00..0260 ation fractio--1150 IIGGIIIVI IGVIII
r z
-11 e0 SWP48111LL ionig -20
1 o
/ l
f λ -25 IGII IGIX
SWP48112HL
(Gaussian 7Ao FWHM)
0
-7 -6 -5 -4 -3 -2 -1 0 1 2
1000 1200 1400 1600 1800
o log m / g/cm2
wavelength / A
Fig. 5 Ionization fraction of the generic iron-group ele-
Fig. 4 TMAP SEDs with two different values of EB−V
(reddeninglaw of Fitzpatrick 1999, RV =3.2) compared to mentinourmodelatmospherewithTeff=55000K,logg =
the FUSE and IUE observations of EC11481−2303. The 5.8, [H]=0.127, [He]=−1.405, [C]=−0.875, [N]=−0.875,
reddening (EB−V = 0.06) and the interstellar Hi column [O]=−0.875, and [IG]=0.000.
density (3·1020cm−2) are adopted from Stysetal. (2000).
All SEDs are normalized to match the observed flux at
the number ratio LIN:POS=475750:1100) results in
1950˚A.TheFUSEandIUEhigh-resolutionobservationsare
an unrealistically high flux level.
convolvedwithaGaussian inordertomatchtheresolution
In Fig.6, we compare theoretical SEDs calculated
of theIUElow-resolution observation.
withdifferentiron-groupabundanceswithobservations.
Atatentimessolarabundance,theagreementwithob-
that the UV flux appears much flatter than predicted servations is good except for 1100˚A <λ<1420˚A. In
by a H+He model. thiswavelengthinterval,thetheoreticallypredictedflux
At the higher Teff derived by us, radiative levita- is still about 20% higher than observed. Two possible
tion could be very efficient. This supports our idea reasonsmaybethattheindividualiron-groupelements
that line blanketing of overabundant iron-group ele- aredifferentlyenhancedbyradiativelevitationandthat
ments is responsible for the flat UV flux. We cal- there are still not enough lines considered.
culated model atmospheres with different iron-group One of the main challenges for the spectral analy-
abundances ([IG]=0, 1, 2). For these calculations, sis of EC11481−2303 is how to include all the lines
the complete iron group is represented by a generic in the TMAP calculations, i.e. where to get all the
model atom which was created by our IrOnIc code data from. Kurucz’ data files were recently extended
(Rauch & Deetjen 2003). Fig.5 shows the ionization (Kurucz 2009) but mainly for the lowest ionization
fractions of the generic iron-group element (IG). IGiv stages. The only exception is Fe where new data files
to IGvi are the dominant ionization stages in the line- are available up to Fevi. For Fev, e.g., the number
forming region in the relevant parameter range. In our of LIN lines increased from 1000385 to 7785320. We
calculations, we consider IGii - ix. calculated a new model with the extended LIN lists.
Our main data source for iron-group elements are Although many more lines are considered on an even
Kurucz’ data files (Kurucz 1997). For line transitions, more refined frequency grid, the differences are within
theseareavailableasso-called“POS”lists,thatinclude a few percent only (Fig.7).
lines with laboratory measured “good” wavelengths, Changes in the individual abundances of the iron-
and “LIN” lists, that include theoretically calculated group elements result in stronger changes in the UV
wavelengths in addition. TMAP uses the LIN lists for flux. We performed a test calculation and increased
model-atmosphere calculations in order to have a reli- theNiabundanceintheabundancesolarpatternofour
abletotalopacity. Foracomparisonwithobservations, generic IG model atom by a factor of ten (Fig.8). The
in general POS line lists are applied. For our compar- flux is reduced in sections of the UV spectrum where
ison of the observed UV flux of EC11481−2303, how- Ni lines dominate.
ever,we have to use LIN lists because due to the much We conclude that fine tuning of all iron-group ele-
lower number of lines in the POS lists (e.g. for Fevi is ments is necessary in order to achieve a better agree-
ment with observations. This is corroborated by a
EC11481−2303 -APeculiarSubdwarfOBStarRevisited 5
o-1A
-1c1 Kurucz 1997
e Kurucz 2009
s
-2m
c
g
r
1 NEBH-VI == 03..0061020 -11 e0
1
/
fλ

0
o-1A 1000 1200 1400 16o00 1800
-1c wavelength / A
e
s Fig. 7 Comparison of two SEDs of our [IG]=1 model,
-2m calculated with theold and theextended LINline lists.
0
c
g
r
e

-1110 o-1A
/ fλ HH++HHee+C+N+O+IG,[IG]=0 -1ec1 [[NNii//FFee]]==01
s
1 -2m
2
c
g
r
1 e

-110
1
/
fλ

0
1000 1200 1400 1600 1800
o
wavelength / A
Fig.8 ComparisonoftwoSEDsofour[IG]=1model(cal-
0
culatedwiththeextendedLINlinelists),withtwodifferent
1000 1200 1400 1600 1800
o Ni/Feabundanceratios, compared with observations.
wavelength / A
Fig. 6 SEDs of our final model atmospheres (parameters
Fig.10comparesatheoretical[IG]=1SEDwiththe
see Fig.5) with different iron-group abundances compared
FUSE observation. Note the excessively large differ-
with observations. At bottom, the SEDs and the FUSE
observation areconvolvedwith aGaussian of5˚A(FWHM) ence in the numbers of POS and LIN lines that are
forclarity. Inthetoppanel,onecangetanimpressionofthe marked at the top and bottom, respectively. As well
extremelyhighnumberofspectrallinesthatareconsidered. as in Fig.9, we see an indication that fine tuning of
iron-group abundances is necessary. Some prominent
absorption features, e.g. at λ1117˚A and λ1122˚A, ap-
comparison of two of our model SEDs, calculated with
pear not reproduced by our model. These might either
[IG]=0 and [IG]=1 and the LIN line lists with the
stem from the iron groupbut then, they are not in the
high-resolutionIUEobservation(Fig.9). ThehigherIG
LIN lists or their wavelengths are uncertain and they
abundances yield a much better agreementwith obser-
appear at a wrong wavelength,or they are of interstel-
vations. Acloselookshowsthatindividuallinescanbe
lar origin(in case of λ1122˚A, Feii is a goodcandidate
identified,e.g. CoVλλ1310.07,1329.80˚A,whichagree
foranidentification). Rauch etal.(2007)didshowthat
well with observations and will allow a precise abun-
a simultaneous fit of both, stellar spectrum as well as
dance determination. On the other hand, some strong
interstellar line spectrum, allows to identify isolated,
features appearinthe synthetic SEDs but they are not
unblended stellar lines and to improve both, the ISM
marked, e.g. around 1328˚A. These are found only in
modelaswellasthephotosphericmodel. Inthefurther
the LIN line lists and their true wavelength positions
course of our EC11481−2303 analysis, we will include
may be far off.
the interstellar lines in our modeling.
6
CoVNiV CrIVNiVFeVNiVNiIVCoVCrVIINiVV IVNiVCoVNiVFeVNiVCoVNiVCrIV CoVNiVCoVNiVFeVNiVCoVNiV CrIVCoVCrIVNiVCrVIIFeVNiVFeVV IV CoVFeVCoVNiVCrIVNiVCoVCrIVNiVCoVCrIVCoV V IVCoVV IVCoV CrIV CoV V IVNiVCrIVCoV
1.0
x
u
e fl
v
ati
el
r

0.5
IUE SWP48112HL
CrIV, loggf=0.794
[IG]=0
[IG]=1
0.0
1310 1312 1314 1316 1318 1320 1322 1324 1326 1328 1330
o
wavelength / A
Fig.9 Comparisonoftwosyntheticspectrawithdifferentiron-groupabundanceswiththehigh-resolutionIUEobservation.
Thestrongestlinetransitions(loggf ≥−1)fromKurucz’sPOSlistsaremarkedatthetop. Theirloggf valuesareindicated
at thebottom by vertical bars.
CrIVNiVCrIVFeVCrIVV IVCrIVV IVCoVFeVV IVFeVV IVCrVFeVCrIVFeVCoVNiIVCrIVTiVFeVCrVCrIVCrVIIFeVCoIVNiVCoIVFeVNiVCrVCrIVCrVIICrVFeVIICrVFeVCrV FeV CoV CrIVFeVNiVFeVCrV FeVFeVCrVFeVCrVFeVCrIVTiVV IVNiIVNiVFeVNiVNiVFeV FeVCrVCrIVCrIVCoIVFeV ScIVCrVCrVV IVNiVFeVFeVTiVCoIVFeVCrIV
1.5
eIIis
F
x
u
e fl1.0
v
ati
el
r

0.5 FUSE
[IG]=1
1112 1114 1116 1118 1120 1122 1124 1126 1128
o
wavelength / A
Fig. 10 Same as Fig.9, but for a [IG]=1 model and the FUSE observation. At the bottom, the loggf values of the
strongest line transitions (loggf ≥−1) from Kurucz’sLIN lists are indicated. “is” denotesinterstellar lines.
EC11481−2303 -APeculiarSubdwarfOBStarRevisited 7
5 Results and conclusions
2 HHe C N O CaScTi VCrMnFeCoNi
The opticalandUV observationofEC11481−2303are
reproduced by our TMAP models with Teff=55000± 1
5000K, logg=5.8 ± 0.3, [H]=0.127, [He]=−1.405,
[C]<−0.875, [N]<−0.875, [O]<−0.875, and [IG]>1. X]
C, N, and O abundances are upper limits only, de- [ 0
termined from the optical spectrum. Our calculations
show that Ni is probably as abundant as Fe. The error
-1
in the H and He abundances is about 0.3dex.
Our suggestionis that the photospheric abundances
of EC11481−2303 (Fig.11) display the interplay of
5 10 15 20 25
gravitational settling and radiative levitation. The atomic number
latter is responsible for the strong iron-group over-
abundance. Fig. 11 Photospheric abundancesof EC11481−2303.
EC11481−2303and other sdOB and sdB stars with
such a high iron-group abundance (UVO0512−08,
PG0909+276, and UVO1758+36, Edelmann 2003)
should be subjects to detailed diffusion calculations. analysis, a precise determination of the reddening (in-
Therearestilldeviationsbetweenthesyntheticspec- cluding infrared measurements like e.g. 2MASS) and
trum and observations (Fig.6, 1100˚A < λ < 1420˚A). the interstellar line absorption will be performed.
Most likely, these are due to our assumption of a so- The extension of Kurucz’s data files is highly desir-
lar abundance pattern within our representation of Ca able because we are strongly hampered in the spectral
– Ni by one generic model atom. A detailed spectral analysisof hot starsby the lack of reliable atomic data
analysiswithindividualmodelatomsforalliron-group of the expected high ionization stages. This is a chal-
elementsisstillon-going. Withintheframeworkofthis lenge for the near future.
Acknowledgements T.R. is supported by the Ger-
manAerospaceCenter(DLR)undergrant05OR0806.
J.W.K. is supported by the FUSE project, funded by
NASA contract NAS532985. This research has made
use of the SIMBAD Astronomical Database, operated
at CDS, Strasbourg, France.
8
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